Bethe-Block. Stopping power of positive muons in copper vs βγ = p/mc. The slight dependence on M at highest energies through T max

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1 Bethe-Block Stopping power of positive muons in copper vs βγ = p/mc. The slight dependence on M at highest energies through T max can be used for PID but typically de/dx depend only on β (given a particle and medium) At low β -de/dx 1/β 2 decreases rapidly as β increases. At relativistic velocities β 1 and reaches a min at βγ 3 (a particle at the energy loss min is called mip). Beyond the min the energy loss increases logarithmically (due to the increase of T max and b max ). However as the range of distant collisions extends, the atoms close to the path of the particle will produce a polarization which results in reducing the electric field strength acting on electrons at large distances Density effect: δ/2 The relativistic rise depends on ln(βγ) but in the ultrarelativistic region only on lnγ hence 8 on the particle mass (used for PID)

2 Bethe-Block For a given particle (z) and target (I,N,Z,A), the energy loss depends only on the velocity of the particle! Most relativistic particles have energy loss rates close to the minimum (mip 9 = minimun ionizing particles)~ 2 MeV/g/cm 2

3 δ rays and restricted energy loss The kinetic energy of the particle is transferred to electrons (δ-rays)( emitted in directions close to the incoming particle one (cos( cosθ 1/ 1/T max ).. Therefore the energy of the particle can result in secondary ionization processes with additional δ-rays traveling far away from the emission point. Often a restricted energy loss is used to estimate the deposited energy in a detector using instead of T max the effective detectable max transferred energy (= effective average max δ-rays energy that can be absorbed inside the device). 10

4 Fluctuations in energy losses: Energy losses of massive charged particles are a statistical phenomenon (collisions are a series of independent events) and in each interaction different amounts of kinetic energy can be transferred to atomic electrons. The energy lost by a particle crossing a path x has an energy distribution called energy struggling function (not for electrons for which the collision process is not dominant) which is the solution of transport equation: The pdf describing the distribution of energy loss Δ in an absorber thickness x is called Landau distribution The most probable energy loss Δ p x(a+ x(a+lnx) ) and the ratio of the full width at half maximum (w = FWHM) w/ Δ p 1/x For very thick absorbers where the energy loss exceeds 1/2 Einitial f(δ) ) approximates a Gaussian 11

5 Range of charged particles in liquid H 2, He gas, Carbon, Iron, lead A proton of 10 GeV M ~1 GeV will cross R = 6000 g/cm 2 in Fe(true if interaction length > R) Range At high energy -de/dx depends on β R depends on βe βe = pc R/M depends on pc/m = βγ Due to the statistical nature of energy loss, mono-energetic particles do not travel the same paths in absorbers (range straggling that follows a Gaussian distribution) 12

6 PID using de/dx The low momentum region where -de/dx 1/β 2 and the relativistic rise depend on m so can be used for PID. While in 1 thin layer of absorber the mean of the energy loss distribution is affected by the long tail of the Landau asymmetric distribution, many consecutive layers can be used to make a precise measurement of energy loss that allows (for known momentum) to determine particle masses. For a fixed layer thickness L/N of N chambers the resolution improves as 1/sqrt(NL) 13

7 Radiation Energy Losses of Electrons and Positrons At low energies electrons and positrons lose energy by ionization and scattering processes (Møller( ller, Bhabba,, e+ annihilation). Ionization energy losses rise logaritmically with energy, while radiative losses rise almost linearly so high energy electrons (E critical ~550/Z in MeV for Z>13) ) and positrons loose energy by radiative emissions (bremsstrahlung when traversing matter and synchrotron radiation when occurring in circular acceleration) ) and e + e - pair production.. Radiation emission is connected to acceleration or deceleration of charged particles.the emitted radiation per unit time depends on the velocity variation: de dt = 2e2 2 dv 3c 3 dt For bremsstrahlung by a particle of charge z and mass m in an absorber of atomic number Z the acceleration depends on zze 2 /m p de and hence it is much less probable for massive particles z2 Z 2 than electrons dt B m 2 The characteristic amount of matter traversed while bremsstrahlung is called radiation length X 0 (in g/cm 2 ): the layer thickness that reduces the mean energy by a factor of e: e and pair productions occur with a characteristic length of L p ~9/7X 0 de E = dx σ bremss (E > m e ) α 3 z 2 4r 2 e ln 183 E = E 0 e X X 0 Z + 1 1/3 8 X 0 716A A X 0 (g / cm 2 ) = +1) ln Z(Z +1.3) ln 183 Z(Z+1.3) accounts for interactions between Z + 1 = N A σ bremss 1/3 1/ high energy electron and those bound to atom

8 γ Energy losses of electrons electron nucleus The process can happen only in the field of an atom or of atomic electrons to conserve energy-momentum ATMOSPHERE 38 g/cm 2 => 280 m => atmosphere is about 25 X 0 at the vertical LEAD 6.37 g/cm 2 / g/cm 3 = 0.56 cm 15

9 Atomic and Nuclear Properties of materials pdg.lbl.gov/2005/reviews/.gov/2005/reviews/atomicrpp.pdf 16

10 Multiple Coulomb Scattering When a charged particle passes in the neighborhood of a nucleus its trajectory is deflected.. The Coulomb scattering distribution is well described in the Moliere Theory. For small scattering angles it is roughly Gaussian,, but at larger angles It has larger tails than a Gaussian. Instead of considering the total deflection θ it is convenient to consider its projection on a plane. In the Gaussian approx (small deflection angles), the root mean square value (rms( rms) ) projected on a plane is θ 0 = θ rms proj = 2 θ proj = 1 2 θ 2 = 1 2 θ rms x,y orthogonal to direction of motion The space and plane angular distributions are approximately given by 2 θ space 2 2θ plane Product of 2 distributions on the plane 1 2πθ exp 2θ 2 plane dθ plane 0 2θ 0 17

11 Muon energy losses Ionization Stochastic losses ~2 MeV/g/cm 2 (dominate > 1TeV ) rock Pair production E dx 1 1 R µ = de de = log(1 + E / E de a + be b 0 E 0 c ) photonuclear ionization E c = a / b critical energy bremsstrahlung 18

12 Interaction of Radiation with Matter Beams of monochromatic photons of initial intensity I 0 are reduced exponentially in intensity I = I 0 exp(-µ att x) µ att =attenuation coefficient = σ N A ρ/a The main processes contributing to the total cross-section are: - Photoelectric effect: interactions with entire atomic electron cloud resulting in complete absorption of photon energy - Thompson and Compton scattering on atomic electrons at γ energies >> electron binding energies and electrons are quasi-free - pair production: the photon incoming energy is high enough to allow e + e - creation in the Coulomb field of an electron or of a nucleus Cross section remains flat up to 10 TeV 19

13 The photoelectric effect 1905 Einstein The phenomenon is responsible of opacity in stellar interiors and atmospheres. At low photon energies hν << m e c 2 it is the dominant process by which photons lose energy. If the energy of the incident photon is hν it can eject electrons with binding energy B e hν. The energy of an electron leaving the atom is K e = hν - B e. If the electron energy is lower than the binding energy of a shell an electron from that shell cannot be emitted. Therefore the absorption curve exhibits absorption edges whenever the incoming photon energy matches the ionization energy of K, L, M, shells (that can have substructures except for K-shell) The photoelectric absorption probability is larger for more tightly bound electrons (K-shell) since free-electrons cannot absorb photons. The binding energy for K shell is B e (K) R y (Z-1) 2 ev where R y =13.61 ev Rydberg constant 20

14 The photoelectric effect The K shell cross section (hν( << mc 2 ) depends strongly on the medium! If an atomic electron is emitted due to the photoelectric effect, a vacancy is created in the shell leaving the atom in an excited state. The atom can readjust itself to a more stable state via a radiative emission (X-ray photon) or 1 or more electrons (Auger effect) σ K = σ Th 4 2 Z 5 mc 2 7/ hν σ Th = ( 8 / 3)πr 2 e = cm 2 21

15 Compton Scattering 1923 Compton discovered that the wavelength of hard X-ray radiation increases when it is scattered from stationary electrons. This is another phenomenon in which radiation appears in its corpuscolar nature (duality wave- particle): scattering of photons on quasi-free. At lower energies Thompson scattering dominates in which the wavelength of photons is not changed. A classical description of the collision leads to a total cross section σ Th = 8π 3 r e 4 e2 = 6πε 2 0 m 2 e c = m 2 The # of electrons/v =N e decreases exponentially with distance dn dx = σ tn e N N = N 0 exp( σ T N e dx T ) mc 2 (ν ν') = hνν'(1 cosθ ν ) hν = hν'+ h 2 optical depth of the medium due to Thompson scattering τ = σ T N e dx reduced photon energy Exercise mc 3 νν'(1 cosθ ε = hν 2 ν ) = hν' [ 1+ ε(1 cosθ ν )] mc 2 Demonstrate the Compton shift formula Δλ = λ λ = λ e (1-cos cosθ ν ) with λ e = h/(mc) Compton wave length of electron 22

16 Compton Scattering Hence the fraction of the incoming photon carried by the scattered photon is hν' hν = 1 1+ ε(1 cosθ ν ) angle. [ ] cosθ ν =1 1 ε hν' hν 1 is the cosine of the photon scattering The photon is forward scattered when hν hν and for low energies (ε << 1) hν hν. Under these circumstances K e 0 The electron kinetic energy is K e = h(ν ν') = hν' [ ε(1 cosθ ν )] = hν ε(1 cosθ ) ν 1+ ε(1 cosθ ν ) The maximum kinetic energy is for backward scattering of photons θ = 180deg K e = hν 2ε 1+ 2ε Fraction of the incident photon energy taken by the Scattered photon 23

17 The cross-section for Compton scattering Klein-Nishina formula For an atom with Z electrons σ = Z σ K-N For low energies ε 0 dσ dω ε >> 0 dσ 1/ε And for dω C Compton int prob decreases the higher the photon energy C dσ dω Th = r 2 e 2 (1+ cos2 θ ν ) The higher the photon energy The more anysotropic the Scattering (forward) ε 0 σ C σ Th = 8 3 πr 2 e ε >>1 σ C 3 8 σ Th 2ln(2ε) +1 ε 24

18 Inverse Compton Scattering Most important process in high energy astrophysics, eg when accelerated electrons collide with MWB radiation or other ambient fields ultrarelativistic electrons scatter low energy photons to high energy. In the reference frame π 0 25